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Outer Planet Magnetospheres: a Tutorial

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Outer Planet Magnetospheres: a Tutorial OUTER PLANET MAGNETOSPHERES: A TUTORIAL C. T. Russell Department of Earth and Space Sciences and Institute of Geophysics and Planetary Physics, University of California Los Angeles, 405 Hilgard Avenue, CA 90095-1567, USA ABSTRACT Outer planetary magnetospheres represent giant laboratories for testing our ideas of how magnetospheres behave. In this tutorial review we examine the role of external and internal pressure in determining the size and shape of the magnetosphere. We examine the relative roles of reconnection with the solar wind magnetic field and the mass addition inside the magnetosphere in driving the circulation of plasma in the magnetosphere. We also examine how the jovian magnetosphere maintains a steady state in an average sense despite the continued addition of mass deep in the magnetosphere. INTRODUCTION The sun’s ionized corona expands rapidly, nearly isotropically, into the surrounding heliosphere carrying the solar magnetic field with it. This expanding magnetized plasma interacts with the planets creating planetary magnetospheres, regions of enhanced magnetic field strength surrounding the planet where magnetic stresses play a dominant role in controlling the flow of the plasma. The nature of the obstacle to the solar wind flow is important in determining the nature of a planetary magnetosphere. If the planet has its own strong magnetic field due to an interior magnetic dynamo or a remanent field of a magnetized crust, then the magnetosphere is termed intrinsic. The size of an intrinsic magnetosphere is determined by the location of the point where the dynamic pressure of the solar wind flow is equal to the magnetic pressure of the intrinsic magnetic field. For Mercury this occurs just above the surface of the planet, whereas for Jupiter it occurs close to 100 times further from the center of the planet than the cloud tops. Because the solar wind flows supersonically, that is the flow is much faster than the speed of a compressional wave that could deflect the plasma flow, then a bow shock forms that slows, heats and deflects the flow before it reaches the magnetic obstacle. This shock stands off in front of the obstacle at a distance sufficiently far that the compressed solar wind plasma can flow around the obstacle. Induced magnetospheres occur when the region of strong magnetic field arises not from magnetism within the planetary obstacle but from the interaction of the solar wind with the obstacle. Induced magnetospheres in turn can be divided into two classes, one in which the planetary body is adding mass to the solar wind flow in the form of ions newly created from a neutral atmosphere. This occurs most spectacularly at comets and less so at Venus and Mars and also at the moons Io and Titan where a “magnetosphere” is created by a flowing planetary wind. The other class of induced magnetospheres arises from classical electromagnetic induction in a highly electrically conducting medium. This conductor can be the electrically conducting interior of a planet or moon such as an iron core or a salt-water ocean, such as in Europa, or it can be the ionosphere of a planet above the surface, such as at Venus and Mars. In these cases the magnetic field external to the planetary body is excluded from the interior of the conductor for a length of time dependent on the size and electrical conductivity of the obstacle. This time can be long compared to the time scale of directional changes in the exterior magnetic field. Currents flow in the conductor to oppose the field change and the field outside obstacle increases, creating a magnetic barrier that in turn deflects the flowing plasma. As in the case of intrinsic magnetospheres, either class of induced magnetospheres can lead to the formation of a bow shock, standing in front of the obstacle when the flow velocity exceeds the velocity of the compressional wave that is required to deflect the flow around the obstacle. Interplanetary Magnetic Field Tail Current Plasma Mantle -1 Magnetic Tail Northern Lobe Polar Cusp Plasma Sheet Plasma- sphere X 2 -1 Field Line Streamline 1 Foreshock Neutral Sheet Current Boundary Field-Aligned Current Ring Current Magnetopause Solar Wind Magnetopause Current Y Fig. 1. Cut away model of the terrestrial Fig. 2. Magnetic field lines and flows in the plane magnetosphere showing the plasma regions, containing the upstream field and flow passing through the magnetic field lines, electric currents and flows. stagnation point as calculated in the convected field gas dynamic model of the interaction of the solar wind with the Earth’s magnetosphere (after Spreiter et al., 1966). Herein we examine the physics of the outer planet magnetospheres. We have visited the magnetospheres of Jupiter and Saturn many times but those of Uranus and Neptune only once. These magnetospheres are all intrinsic. No spacecraft has yet visited Pluto, but unless the atmosphere and interior are completely frozen we would expect one of the classes of magnetospheres to form. Above we discussed briefly the size of intrinsic planetary magnetospheres as determined by a balance of pressures, between the dynamic pressure of the flowing solar wind and the magnetic pressure of the intrinsic magnetic field. Figure 1 shows an intrinsic magnetosphere as deduced from observations of the terrestrial magnetosphere. The magnetosphere forms a cavity in the flowing solar wind and throughout much of this cavity the pressure is dominated by the magnetic field. The pressure that determines the size of the cavity is normal to the surface of the cavity. Outside the cavity is plasma whose pressure decreases from outside to inside across the cavity boundary. This pressure gradient exerts an inward force. It is balanced by a magnetic pressure force pushing outward as the magnetic field increases from outside to inside across the cavity boundary, or magnetopause. The relationship between this normal component of the pressure and the directed dynamic pressure in the upstream solar wind is complex and is only understood semi-empirically (e.g. Petrinec and Russell, 1997). Similarly, the magnetic field depends on a multitude of currents, not just those interior to the planet. Thus the point of pressure equilibrium or force balance can move. Since the magnetic flux crossing the surface of the planet is fixed on the time scale of most magnetospheric “events”, any variations in the magnetopause location under conditions of constant external pressure must involve a change in shape of the magnetosphere. Rapidly rotating magnetospheres with strong internal sources of plasma add an additional complexity that will be discussed below. The magnetosphere illustrated in Figure 1 is stretched in the antisolar direction. This magnetotail, as it is called, is a region in which energy can be stored, analogous to the energy that can be stored in an electrical circuit by an inductor. However, to understand the behavior of the magnetosphere one has to understand the transport of mass and magnetic flux that occur on time scales much slower than the speed of light. Thus one can often catch a magnetosphere in the act of changing, or in a metastable state, ready to change. One can add energy to this system by simply compressing it, just as one can store energy in a compressed spring. The work done, or energy stored, is the force times the distance moved normal to the surface, integrated over the magnetosphere. One can also add energy to the system by eroding the magnetic field on the dayside and 2 Me V E Ma J S U 2.2 Jovian Magnetopause Subsolar Standoff Voyager 1 & 2 Z JSM < 20 R J 10 β 1.0 J 2.0 Slope = -0.22 Beta N 1.8 Mms 10 Log (R ) [R ] 1.6 Fast Magnetosonic Mach Number 1.4 1.0 0.1 0.1 1.0 10 -2.0 -1.5 -1.0 -0.5 0 Log 10 (Psw) [nPa] Heliocentric Distance (AU) Fig. 3. Expected heliocentric variation of the fast Fig. 4. The location of the magnetopause nose plotted magnetosonic Mach number of the solar wind flow versus the solar wind dynamic pressure for Voyager 1 relative to the planets and the solar wind beta or ratio of and 2 data. The line gives the best fit to the data on a the magnetic to thermal pressure (after Russell et al., log-log scale (after Huddleston et al., 1998a). 1982). carrying it into the tail. This requires a tangential stress such as friction at the surface of the magnetosphere. This tangential stress can be a viscous interaction associated with waves on the boundary in the presence of dissipation, with particles being scattered from the flowing, shocked solar-wind plasma into the magnetosphere, or by magnetic coupling of the magnetospheric and solar wind magnetic field. This coupling is illustrated in Figure 1 above the magnetosphere and behind the depression in the shape of the magnetopause called the polar cusp. Here the magnetic field lines are bent in the direction as to slow the solar wind flow. As the solar wind slows in this region, mechanical energy is removed from the flow, and is stored in the tail lobes where the magnetic field energy increases. The transport of energy into the tail occurs through an electromagnetic Poynting flux. Magnetospheres are vast in scale but contain very little total mass. The stress applied to the magnetosphere by the solar wind must be ultimately taken up by the planet. One of the ways to transmit this stress to the ionosphere and thence through collisions to the upper atmosphere and the planet, is through currents parallel to the magnetic field as illustrated in Figure 1. These currents close on pressure gradients in the magnetosphere. Another is through the gradient in the Chapman-Ferraro currents acting on the planet’s dipole (Siscoe, 1966).
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